Wavefront Sensor for Millimeter/Submillimeter-Wave Adaptive Optics Based on Aperture-Plane Interferometry
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Wavefront sensor for millimeter/submillimeter-wave adaptive optics based on aperture-plane interferometry Yoichi Tamuraa, Ryohei Kawabeb, Yuhei Fukasakuc, Kimihiro Kimuraa, Tetsutaro Uedaa, Akio Taniguchia, Nozomi Okadad, Hideo Ogawad, Ikumi Hashimotoe, Tetsuhiro Minamidanib,f, Noriyuki Kawaguchib, Nario Kunoc, Yohei Togamia, Masato Hagimotoa, Satoya Nakanoa, Keiichi Matsudaa, Sachiko Okumurag, Tomoko Nakamurag, Mikio Kuritah, Tatsuya Takekoshij,i, Tai Oshimab, Toshikazu Onishid, and Kotaro Kohnoi aNagoya University, Nagoya, 464-8602 Japan bNational Astronomical Observatory of Japan, Mitaka, Tokyo, 181-8588 Japan cUniversity of Tsukuba, Tsukuba, Ibaraki, 305-8573 Japan dDept. Physical Science, Osaka Prefecture University, Sakai, Osaka, 599-8531 Japan eDept. Aerospace Engineering, Osaka Prefecture University, Sakai, Osaka, 599-8531 Japan fNobeyama Radio Observatory, Nagano, 384-1305 Japan gJapan Women's University, Bunkyo, Tokyo, 112-8681 Japan hKyoto University, Kyoto, 606-8501 Japan iInstitute of Astronomy, The University of Tokyo, Mitaka, Tokyo, 181-0015 Japan jKitami Institute of Technology, Kitami, Hokkaido, 090-8507 Japan ABSTRACT We present a concept of a millimeter wavefront sensor that allows real-time sensing of the surface of a ground- based millimeter/submillimeter telescope. It is becoming important for ground-based millimeter/submillimeter astronomy to make telescopes larger with keeping their surface accurate. To establish `millimetric adaptive optics (MAO)' that instantaneously corrects the wavefront degradation induced by deformation of telescope optics, our wavefront sensor based on radio interferometry measures changes in excess path lengths from characteristic positions on the primary mirror surface to the focal plane. This plays a fundamental role in planed 50-m class submillimeter telescopes such as LST and AtLAST. Keywords: Submillimeter, single-dish telescope, adaptive optics, aperture-plane interferometry 1. INTRODUCTION The increase in size of telescopes, i.e., the acquisition of larger collecting area and spatial resolution, is the very history of observational astronomy. For millimeter and submillimeter-wave telescopes (mm/submm telescopes, hereafter), increasing the antenna diameter with keeping the mirror surface accurate is important in developing new astronomical fields. The advent of an aperture synthesis interferometer has resulted in significant improve- ment in collecting area and angular resolution. For example, the Atacama Large Millimeter/submillimeter Array arXiv:2102.09286v1 [astro-ph.IM] 18 Feb 2021 (ALMA) realized the performance by keeping the size of the array element antennas moderate (D = 7 and 12 m) and increasing stiffness of them to keep their surface down to 20 µm r.m.s or even better. However, aperture synthesis interferometry is not a panacea in realizing a larger-scale mm/submm telescope. The focal plane instruments of radio interferometers are limited to coherent receivers (e.g., heterodyne receivers) that can detect the phase of celestial signals. On the other hand, large-format arrays of direct photon detectors (e.g., cameras and integrated superconducting spectrometers1), which have prospered in recent years, cannot be accommodated as receivers for an aperture synthesis interferometer, because phase detection is not possible. Further author information: (Send correspondence to Y.T.) Y.T.: E-mail: [email protected], Telephone: +81 (0)52 789 2846 1 Thus, it becomes increasingly important to have a large-aperture single-dish mm/submm telescope (e.g., Large Submillimeter Telescope2 and Atacama Large Aperture Submillimeter Telescope3), especially in the field in which direct photon detector arrays play a crucial role. A limiting factor in size and operating frequency of mm/submm telescopes is the deterioration of optical performance due to changes in the environment surrounding the telescope, such as wind load and thermal deformation in addition to gravitational deformation.∗ This is unlike the visible/near-infrared (NIR) regime, where the major cause of wavefront degradation is due to temperature inhomogeneities of the tropospheric atmosphere on the order of tens of cm. Mm/submm telescopes have been manufactured on the premise of large (> 10 m) structures in order to ensure the large collecting area and high spatial resolution. For this reason, building a huge dome to cover it would require enormous costs. Therefore, it was an issue to construct an `exposed' antenna outside and ensure its mirror surface with an accuracy of ∼ 20 µm. With the existing technology, it is possible to measure the mirror surface shape in advance by the photogram- metry or radio holography methods, and to correct the mirror surface down to the level of ∼ 10 µm (r.m.s.) by adjusting the positions of the primary mirror panels. In recent years, the weight reduction and enhanced stiffness of the primary mirror structure and the active surface control of the primary mirror panels by the motorized actuators have been used to keep the primary mirror an ideal shape with respect to changes in telescope elevation angle.4 The problem is, however, the deformation of the primary mirror surface and the secondary mirror support structure due to wind load and thermal deformation. This is because the time-scale of the deformation is fairly short compared with a typical time-scale of astronomical observations. It is difficult to measure the mirror surface shape in real time during the observations. It is likely that most of the deformation across the primary mirror are dominated by lower-order deformation modes. Their typical spatial scale is found to be a fraction of the aperture diameter (1{10 m).5 The time-scale of the deformation is determined by the natural frequency of the primary mirror structure to the wind load and is typically ∼ 10−1{1 s (e.g., ' 1:1 s for the Nobeyama 45 m mm-wave telescope6). In this case, the control technology of the adaptive optics does not become a major obstacle. The spatial and time-scales of adaptive optics control, which is already established in NIR astronomy, are on the order of ∼ 10 cm and ∼ 1 kHz, respectively, which are several orders of magnitude higher in spatial and temporal frequencies than those of large mm/submm telescopes. Therefore, in the radio telescope, the wavefront compensation can be performed at the primary or secondary mirror equipped with the active surface control, instead of a high-speed deformable mirror employed in NIR telescopes. Then, the major obstacle in the mm/submm regime lies in how to measure the wavefront in real time. Wavefront sensing using a Shack-Hartmann sensor is the mainstream in the optical and NIR. A similar approach, however, is not applicable to MAO, since no large-format detector array that can be manufactured cheaply is available in the mm/submm. Instead, radio astronomy can exploit radio interferometry as a native wavefront sensing technology, which measures the difference between the arrival times of the wavefronts coming through two independent optical paths. In this paper, we present a concept of a mm wavefront sensor that allows real-time sensing of the primary mirror surface of a ground-based large-aperture mm/submm telescope. 2. APERTURE-PLANE INTERFEROMETRY FOR WAVEFRONT SENSING As seen in the previous section, what needs to be measured first is the relative time variation of the deformation of the optical system, rather than wavefronts of celestial signals coming through the atmosphere. This is equivalent to the errors in excess path length (EPL) from the focus to arbitrary points on the primary surface. Therefore, the wavefront sensor employs multiple reference microwave/mm sources placed on the aperture of the telescope as a phase standard. The goal of the surface accuracy we measure with a prototype sensor (see later sections) ∗In the millimeter and submillimeter regime, the inhomogeneities in column density of water vapor in the troposphere also cause wavefront degradation. However, the spatial scale at which the water vapor fluctuations dominate the wavefront 2 degradation is typically much longer than the telescope's aperture, for example, & 10 m at the ALMA site. For this reason, in most cases, the Earth's atmosphere is not a major source of wavefront degradation for single-dish mm/submm telescopes. 2 Secondary mirror REFERENCE REFERENCE Primary mirror Delay SWITCH REFERENCE CORRELATOR CORRELATOR REFERENCE CORRELATOR CORRELATOR (a) (b) (c) (d) Figure 1. A schematic diagram of aperture-plane interferometry. (a) A simple interferometer with a correlator which measures a difference in arrival time between two signals coming from a common reference source. (b) The same as (a) but one of the optical paths coming through the free space as shown with the dashed line. (c) The same as (b) but the path through the free space goes through the telescope optics. The lengths of the other path is adjusted by a delay line shown as a loop. (d) The same as (c) with a switch followed by an array of radiators which enables to measure the distances from arbitrary positions to the focal point. is 40 µm r.m.s. This is close to ≈ 100 µm r.m.s., which is routinely achieved by holographic measurements of the Nobeyama 45 m mm-wave telescope. This measurement accuracy is equivalent to the phase accuracy of 1◦ r.m.s. when operating the wavefront sensor at 20 GHz. Figure1 shows the schematic diagram representing the principle of aperture-plane interferometry. A simple interferometer comprising a correlator (figure1a, b) allows us to measure a difference in arrival time between two signals generated by a common reference source. The interferometer in which one of the paths goes through the telescope optics (figure1c) measures the EPL from a certain point on the primary surface to the focus of the telescope. This is similar to what is used for phase calibration (the p-cal method, REF) of phase-referencing receivers used in VLBI Exploration of Radio Astronomy (VERA7). The method is a scalable technique and can be time-multiplexed by an intervening switch which is followed by a series of radiators placed across the surface of the primary mirror (figure1d).